Multiphase Gas Interactions on Subarcsec Scales in the Shocked Intergalactic Medium of Stephan's Quintet with JWST and ALMA

The first description of the NIRCam and MIRI ERO observations, and the choice of Stephan's Quintet as a target (PID 2732), is presented in Pontoppidan et al. (2022). The NIRCam images were taken in six filters (F090W/F277W, F150W/F356W, F200W/F444W), with an integration time of approximately 20 minutes for each band. The FULLBOX five-point "5TIGHT" dither pattern was used, resulting in a large rectangle mosaic of three pairs of dithered tiles, covering 6.3 × 7.3 arcmin². The MIRI image covers a much smaller field of view than the NIRCam mosaic, i.e., the central galaxies, NGC 7318a/b, NGC 7319 (a Seyfert 2 galaxy), and NGC 7320 (foreground galaxy), using four tiles in three MIRI bands, F770W, F1000W, and F1500W. The integration times were about 22 minutes per filter. The large cycling eight-point dither pattern was used to maximize the spatial coverage of a single tile.

We worked with the level 2b images retrieved from MAST. All the HST and JWST data used in this paper can be found in MAST: doi:10.17909/dfsd-8n65.

Observations of CO(2–1) in the three fields of Stephan's Quintet were made with the ALMA 12 m array in a mosaic observing mode on 2015 July 21 (ID: 2015.1.000241.S; PI: P. Guillard). The total integration time for the mosaic was 35 minutes. The spectral setup consisted of four spectral windows of 1.875 GHz with 3.9 MHz channels. One of the spectral windows was centered on the redshifted CO(2–1) (ν_rest = 230.538 GHz), while the remaining three spectral windows covered line-free continuum emission. The standard ALMA calibration plan included bandpass, phase, and flux calibration.

The data were calibrated with the scriptForPI.py calibration script that was included with the archival data, using the ALMA calibration pipeline version r32491 that is included with the Common Astronomy Software Applications (CASA) version 4.3.1 (CASA Team et al. 2022). After calibration, the continuum emission was subtracted in the (u,v) domain by fitting a straight line to the line-free channels. The line data were subsequently imaged using the CLEAN algorithm to produce data cubes with 20 km s⁻¹ channels and a resolution of 036 × 021 (PA ∼ 10°). The rms noise in these data cubes is 0.5 mJy beam⁻¹ chan⁻¹.

A primary beam correction was applied to produce accurate flux measurements. The half-power beamwidth of the primary beam is 228 (10.4 kpc). Total intensity maps of the CO(2–1) emission were made by integrating the signal across the channels where line emission was detected.

Observations of Stephan's Quintet were made with CARMA on 2010 August 1 (program c0593; PI: K. Alatalo) using 15 antennae in the CO(1–0) line. The original data were processed with MIRIAD (Sault et al. 1995). Data reduction and imaging were done in identical fashion to Alatalo et al. (2013). The original data were cleaned (Högbom 1974), which resulted in a restored synthesized beam of 4.1 × 3.3 arcsec² (1.9 × 1.5 kpc). Velocities from extracted spectra have been converted to heliocentric velocities assuming a shift of −11 km s⁻¹ from kinematic local standard of rest to heliocentric. Velocities are all quoted using the optical definition of velocity.

We dearchived calibrated images from the HST MAST archive for filters F665W and F336W with the WFC3/UV instrument. For F665N, the image encompasses the Hα line and includes data based on five dithered observations, each taken with integration times of 5200 s each in 2009 July and August, and obtained as part of a WFC3 ERO campaign (SM4/ERO). For F336W, the observations were obtained as part of proposal ID 12301 (PI: S. Gallagher) on 2011 October 27, with a total integration time of 14,474 s. The resolution of the WFC3/UV observations is comparable with that achieved by NIRCam (005–006). For F665W and F336W, the expected FWHM of a point source is 007, and 0075, respectively, corresponding to ∼30 pc at D = 94 Mpc.

The observations were obtained from the JWST and HST MAST archive, and most images required small adjustments to the WCS coordinates for proper alignment. We initially used Gaia Data Release 3 (DR3) stars to align one image (the HST WFC3 F665N) to the DR3 system using the STScI software package WCSTweak. The NIRCam images were also similarly aligned. However, for MIRI the same method did not produce good results, probably because of the smaller number of stars used for alignment. We used a refinement scheme to solve this problem as we need good alignment between all the images. This was especially important for one of the target fields, which required continuum subtraction of the F1000W and F770W images with F1500W. First, we created small subimages of the F1000W, F770W, and F1500W images around the Field 6 area. The subimage contained at least three Gaia DR3 stars close to the target of interest. These subimages were then aligned more carefully by making small pixel-level shifts in each image to bring the subimages into close alignment with each other and the positions of the Gaia stars. We applied the same method to Fields 5 and 4 (see Section 4). The final results produced local MIRI images which aligned to better than 1 MIRI pixel (03) over the small scale of the extracted regions. The results were validated where the CO ALMA images aligned with clearly related features.

Following Papadopoulos et al. (2008), the uncertainty in the absolute astrometry of the ALMA data is δ θ_bas = (δ B · Δ k )/ B ≈ (δ ϕ_bas/2π)〈Θ_beam〉. The phase error δ ϕ_bas ∼ (2π/λ)(δ B · δ k ), with ∣δ B ∣ ∼ 0.1 mm being the assumed typical error in the baseline length, ∣δ k ∣ = 4.26° the distance to the phase calibrator J2216+3518, λ ∼ 1.33 mm the wavelength of the redshifted CO(2–1) emission, and Θ_beam ∼ 036 the major axis of the synthesized beam. This results in δ θ_bas ∼ 002, which means that the uncertainty in the absolute astrometry of the ALMA data is small compared to that of the HST and JWST data. ¹⁹

Under conditions found in the disks of normal galaxies, strong 7.7 μm and 8.6 μm PAH features are expected to dominate the MIRI F770W filter for low-redshift galaxies (e.g., Smith et al. 2007). In Stephan's Quintet (Appleton et al. 2006, 2017; Guillard et al. 2009; Cluver et al. 2010) much of the mid-IR spectrum of the IGM is dominated by strong pure rotational H₂ lines and emission from [Si ii]λ34.8 μm. Unlike normal star formation regions, the majority of the main IGM filament in Stephan's Quintet exhibits very weak 7.7 μm PAH complex emission and weak nebular lines like [S iii]λ18.7, 33.5 μm (Cluver et al. 2010). This is the result of low star formation rates in the main filament (measured to be between 0.05 and 0.08 M_⊙ yr⁻¹; Cluver et al. 2010), and emission lines more typical of fast shocks than H ii regions (e.g., Konstantopoulos et al. 2014). Examples of these H₂-dominated spectra have been presented in Figures 8, 13, and 14 of Cluver et al. (2010) and Figures 15 and 16 of Appleton et al. (2017). The Spitzer IRS spectra showed that only a small minority of places along the main filament (including parts of the SQ-A region) are spectra found with line flux ratios more typical of star-forming regions.

In Figures 3(a)–(c), we compare the filter transmission bandpass for the MIRI imaging bands (F770W, F1000W, and F1500W) directly with the Spitzer IRS spectra taken within Fields 4, 5, and 6 (Figure 1). The MIRI filter F770W is likely to contain emission primarily from the 7.7 μm PAH complex in regions where star formation dominates, but in other regions, especially those with strong H₂ emission, the 0–0 S(5) can also potentially contribute, as in Figure 3(a). The F1000W filter, on the other hand, almost exclusively contains line emission from the S(3) mid-IR pure rotational H₂ line (see also Figures 15(e) and (h) of Appleton et al. 2017). Based on the presented spectra, we expect the MIRI F1500W filter to be dominated by faint dust continuum from the general IGM, as well as from warm dust from embedded star clusters and sometimes the extended wing of the broad 17 μm PAH complex (Smith et al. 2007); see center panel of Figure 3.

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A clean separation between the mainly PAH-dominated emission in the F770W MIRI filter and the 0–0 S(3) H₂-dominated F1000W band can be seen in many regions along the main IGM filament and in the bridge region in Figure 4. Here, we overlay the contour maps obtained by isolating a PAH band ²⁰ from the IRS spectral map from Spitzer in Figure 4(a) (white contours) with a two-color image of the JWST MIRI F1000W (green) and the MIRI F770W (orange). In Figure 4(b), we overlay the 0–0 S(3) H₂ contours (red) over the same image. The overlays show that the white contours of PAH emission from Spitzer follow quite closely the brighter emission in the F770W-dominated regions, whereas the red contours of H₂ follow closely the green emission originating mainly from the F1000W filter. This figure demonstrates how the 10 μm MIRI filter, as we expected from the individual IRS spectra, detects preferentially H₂ emission, while the F770W filter detects those regions dominated by PAH emission. As shown in Cluver et al. (2010), the PAH emission is relatively faint in the main shock, but this figure emphasizes those regions with stronger PAH emission. This near isolation of the H₂ emission in many regions of the IGM is a result of the relatively narrow width of the F1000W filter and the lack of other lines or bands that can seriously contaminate it for the redshift of Stephan's Quintet.

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Observations using the F1500W MIRI filter are expected to detect mainly warm dust emission, similar to that mapped by Spitzer at 24 μm with MIPS (see Guillard et al. 2010). However, in the SQ-A region (Figure 3(b)), the dust continuum is also contaminated by [Ne iii] emission. As noted above, the very strong 0–0 S(1) line is, fortuitously, redshifted out of the F1500W band.

Before discussing results for each of the zoomed-in regions in the next section, we mention the removal of faint dust emission from F1000W and F770W at the center of Field 6 (near the center of the main IGM filament in Stephan's Quintet). In this field, faint dust continuum in the F1500W filter was seen in the vicinity of the structures we are interested in (see Section 4). We therefore removed the faint dust continuum from both the F1000W and F770W images, using the Spitzer IRS spectrum of the region as a guide. First, we created small subimages of the field in all three bands centered on the main structure of interest and corrected them for small offsets in the WCS, as discussed in Section 2. After removing a median sky background from each subimage, we then subtracted the F1500W image from both the F770W and F1000W images, scaling the continuum by a factor of 0.9 to account for a slight rise in continuum seen between 7.7, 10, and 15 μm in the IRS spectrum of the region. This allowed us to better define the distribution of the H₂ and PAH features. The effect was quite small (the dust emission was quite weak compared with the emission seen in the 7.7 and 10 μm bands) on the resulting fluxes of the F1000W and F770W band. Without spectroscopy of the region on the same spatial scale as the observed features, this method is necessarily imperfect, and where we estimate the fluxes for the features observed in Field 6 in the next section we include larger uncertainties to take this into account. For the subregion near the center of Field 5, no obvious dust structure was seen in the F1500W band near the structures of interest, and so we did not attempt to continuum subtract the H₂ and PAH-dominated emission from the F1000W and F770W images. For Field 4, centered in the SQ-A star-forming region, our attempts at subtracting a scaled version of the dust continuum led to a severe oversubtraction of the emission in the F770W and F1000W bands for most reasonable dust scale factors. This is likely due to a combination of factors. First at that position (see Figure 3(b)), we have shown that the 15 μm band shows emission from not only dust continuum but also relatively strong [Ne iii], as well as broad PAH-band emission from the 17 μm "plateau" complex (see Smith et al. 2007). These additional contaminants make it difficult to isolate the warm dust component in the 15 μm image. Therefore, for this star-formation-dominated field, we conclude that the F1000W and F770W MIRI filters are hopelessly contaminated by dust emission in a way that cannot easily be corrected without high-resolution spectroscopic data from the MIRI Medium Resolution Spectrometer (MRS). This significantly limits our discussion of the molecular hydrogen properties in Field 4. For that region, we limit our discussion of the warm H₂ properties determined from the Spitzer IRS observations.

We also present images in this paper obtained through the NIRCam F200W filter. F150W provides a near-IR continuum band relatively clear of emission lines at the redshift of Stephan's Quintet (0.022 < z < 0.025), with the well-known near-IR [Fe ii] lines falling just blueward or redward of the filter transmission. The bandpass of the F200W filter is sensitive to emission from the hydrogen recombination line Paα and starlight. Diffuse emission from this line becomes obvious when we compare some of these images with those obtained in the F150W NIRCam filter, and by comparison with the HST Hα images.

JWST/MIRI, which is able to detect warm H₂ through the pure rotational transitions, can significantly improve our knowledge of how energy is dissipated in the IGM down to the 03 scales of giant molecular cloud (GMC) complexes (∼140 pc). NIRCam, with its ability to probe down to scales of 005–006 (23–27 pc at 1.5 and 2 μm, respectively) can probe ionized gas and stellar associations and clusters at scales at which sporadic star formation along the main filament is observed.

In this section, we explore a comparison of archival HST and JWST MIRI data with the ALMA CO(3–2).

4.1.1. Field 6: Within the Main Shock-dominated North/South Filament

In Field 6 (Figure 5(a)), we show the integrated CO(2–1) emission from cold molecular gas superimposed on the warm H₂ emission. The warm H₂ gas (assumed to dominate the F1000W filter from the 0–0 S(3) H₂ line) is distributed in an elongated structure extending over 4''–5'' (∼2 kpc) with a triangular-shaped hot spot to the southwest and a series of fainter clumps extending to the east. The overall distribution of warm H₂ has the appearance of a head-tail structure. The CO emission (white contours) displays a clumpy, narrow shell-like structure associated with the warm H₂ head, and a scattering of CO clumps along the tail extending over the same overall extent but not correlating in detail with the hot spots in the warm H₂.

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Based on the published IRS spectra, Figure 5(b) is interpreted as a mix of H₂ (0–0 S(5)) and weak PAH emission. The dominant triangular-shaped head structure seen in the 10 μm image is also strongly represented here. A second major clump of 7.7 μm emission is seen to the northwest, which is weak at 10 μm. It is plausible that this feature is PAH dominated, whereas the emission from the compact head (which is seen also at 10 μm) may be dominated by warmer than normal H₂, which would emit strongly in the 0–0 S(5) line. Without mid-IR spectroscopy, this cannot be verified.

This distribution of ionized gas can be inferred from the narrowband HST F665N filter images (dominated by Hα emission) in Figure 5(c) and the F200W NIRCam image (dominated by Paα emission and near-IR starlight) in Figure 5(d).

Both images show significant extended (presumed ionized gas) emission associated with the head of the warm H₂ including a protrusion, or "spike," labeled in Figure 5(c). The feature is also seen in CO emission. Extended (ionized) gas also follows more closely the distribution of the warm H₂ in the tail than it does the CO emission. This is well demonstrated in the composite image Figure 5(f), which shows an RGB representation of the 10 μm (warm H₂), CO emission (cold H₂), and F200W (ionized gas and stars) images. Gemini optical spectroscopy (Konstantopoulos et al. 2014), which covered this region at a spatial resolution of 1'' (458 pc), strongly suggested that the ionized gas is shock-excited, which would explain why the ionized gas follows the warm H₂ emission in the head and the clumpy tail, and why it is less correlated with the CO emission.

The spatial resolution of the F200W image is far superior to the HST image, and it is clear that there are several compact sources on scales of <005–06 (23–27 pc) embedded in the ionized gas component. These sources are likely stellar associations or unresolved super star clusters (SSCs; Gallagher et al. 2001). They are generally rather sparsely distributed and poorly correlated with the warm or cool molecular hydrogen.

One exception is a group of compact sources and extended F200W emission associated with the southern tip of the northwest clump of CO emission identified in Figure 5(c). These sources may be evidence of recent star formation associated with the head-tail structure. Evidence of this comes from the bright u-band sources associated with the same region shown in Figure 5(e). Although a full analysis of the NIRCam colors of the point sources in Stephan's Quintet will be discussed in a separate paper, we know these blue regions have the optical colors of young clusters with ages estimated to be between 3 and 5 Myr based on previous HST UBV_mVI photometry (Fedotov 2014; see Table B1 for more details). The existence of young star formation near the northwest CO clump may also explain why that region exhibits PAH emission, since PAHs are believed to be excited by UV radiation from young stellar associations.

The kinematics of the diffuse ionized gas centered on the head from optical spectroscopy shows extremely broad line widths (>800–1200 km s⁻¹) in several key emission lines, e.g., [O iii]5007, [O i]6300, and Hα (Konstantopoulos et al. 2014; Guillard et al. 2022). The broad line widths have been attributed to a combination of turbulence and large-scale motions associated with gas along the line of sight. Even broader wings in the Lyα profile were also seen at this position. We will argue in Section 5 that the warm H₂ emission detected by JWST at this position may be responsible for UV scattering of Lyα emission within the turbulent regions and, if so, should exhibit broad line widths.

While we cannot yet measure the kinematics of the warm gas, the CO observations provide kinematic information about the cold molecular phase. Figures 6(a) and (b) show the integrated CO surface density (moment 0) map and the mean velocity field (moment 1) map of the CO emission, respectively. The systemic velocity of the CO emission falls between the velocity of the intruder galaxy NGC 7318b (V(int) = 5770 km s⁻¹) and that of the velocity barycenter of the main group (V(group) = 6600 km s⁻¹). This is consistent with gas which has been decelerated in the collision of the intruder with the group-wide gas. However, the line width of the entire region in CO is less that 120 km s⁻¹, as shown in the integrated ALMA CO(2–1) spectrum of the entire region and the CARMA CO(1–0) spectrum shown in Figure 6(c). This width is significantly smaller than that seen in the ionized gas, implying that the cold molecular gas does not take part in the turbulent motions seen in the ionized gas phase. More extensive kinematic information for the CO emission is presented in Appendix A, where we show individual spectra extracted from many regions (Figure A1), along with tabulated properties of the CO emission on the scale of the ALMA beam (Table A1). The tip of the head structure has the highest radial velocity, and the clumps show a trend to lower velocities as one moves east into the tail, reaching the lowest velocity in one of the most easterly clumps. We will argue later that this suggests material is being stripped from the head into the tail through ram pressure stripping with a hot medium.

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The velocity dispersion ²¹ of the individual clumps around the structure show a mix of unusually broad line widths (up to 100 km s⁻¹ FWHM) and narrow lines (<40 km s⁻¹). In Table A1, we show that the region near the "spike" feature seen in the ionized gas is unusually broad (FWHM =102 km s⁻¹). Away from the head, several regions in the tail of the structure show line widths ranging from 60 to 90 km s⁻¹ (FWHM; see Figure A1) to values lower than 40 km s⁻¹. Given the masses of the individual molecular clumps of a ∼few ×10⁶ M_⊙ (Table A1), it is unlikely that the clumps with high-velocity dispersion are gravitationally bound. ²² The broader lines are consistent with turbulent motions on the subarcsec (∼100–150 pc) scale in the colder gas component, while others have much narrower lines.

One such region occurs in the CO emission clump closest to the bright blue star clusters described earlier. Here, the CO emission is essentially unresolved (40 km s⁻¹ or less), perhaps suggesting turbulent motions are calmer there. Table A1 presents information about the line fluxes and estimated H₂ masses for all the regions measured in detail with ALMA.

4.1.2. Field 5: A Shocked Structure in the Bridge Region

Field 5 lies outside the main molecular filament in Stephan's Quintet, and is closer to the large face-on galaxy NGC 7319. It forms part of an apparent bridge of H₂ emission discussed by Cluver et al. (2010). The JWST images (e.g., Figure 4) show that it is the most easterly of a series of irregular shock-dominated (appearing green in that image) clouds that run approximately east–west across the field and may not be physically connected. In contrast to Field 6, which contained a single coherent CO structure, Field 5 contains three separate bright structures visible in the CO maps (labeled in Figure 7(a)). These consist of a clumpy CO ring about 1'' across (458 pc), a compact elongated core 15 to the southeast of the ring, and a broken linear filament of gas running north–south 05 further east of the compact core. The direction of the elongation of the core points toward the center of the ring. Both Figures 7(a) and (b) show that the compact CO core lies near the brightest emission at the center of the arrow-shaped 10 μm (warm H₂) structure. As the arrow-shaped emission narrows down toward the northwest, it connects to the northern part of the CO emission from the ring. Figure 7(c) shows diffuse Hα emission from the compact CO core but little emission from the ring. The higher spatial resolution of the NIRCam F200W image shows that the corresponding Paα in the core is also elongated Figure 7(d). Faint F200W emission is also seen in the brightest CO ring clump. The U-band image, Figure 7(e), shows no obvious emission near any of the CO molecular structures. Although not shown here, the F150W NIRCam image shows point-like sources in the northern part of the CO ring, suggesting that the F200W image is revealing a mix of H ii regions and compact sources within the ring. The composite image, showing the three gas phases in one figure, emphasizes the warm H₂ connection between the ring and the core (Figure 7(f)).

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A string of clumpy sources at 2 μm is clearly seen approximately 1'' to the east and north of the ring (F200W image). It is not clear if these sources have any real connection to the CO structures, as some could be background galaxies. One of the clumps has a faint U-band association (see Appendix B and Table A3), suggesting recent star formation. Near-IR spectroscopy may help to determine if they are physically associated with Stephan's Quintet by providing evidence of H ii-region-like emission lines at the Quintet's redshift, such as Paα.

The kinematics of CO gas in the compact core is strikingly different from that of the ring and the linear structure. Figures 8(a) and (b) show the integrated and radial-velocity map of the CO(2–1) emission, and spectra of the three different CO-emitting structures are presented in Figure 8(c). The velocities within the partial ring increase clockwise around the ring from 6631 km s⁻¹ in the north, reaching the highest value in the south and east quadrant of 6688 km s⁻¹. In contrast, the compact core (which lies near the peak in the warm H₂ emission) has a much lower (negative 250 km s⁻¹) radial velocity than the ring (∼6400 km s⁻¹). The core also shows a velocity shear along its length of about 30 km s⁻¹ over 05 (200 pc), as shown in Appendix A. Finally, the third linear structure to the east of the compact core shares velocities similar to that of the ring, with velocities ranging from 6620 km s⁻¹ in the south to 6680 km s⁻¹ in the north.

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An in-depth view of the kinematics of Field 5 in given in Appendix A, including individual spectra, Figure A2, and a table of CO properties, Table A2. The spectra show that individual CO clumps in all of the three CO structures exhibit broad CO line widths at the scale of the ALMA beam (∼150 pc). The compact core itself has a velocity dispersion of 145 km s⁻¹ (some of this may be due to the gradient seen along the core), and all components of the ring have line widths which lie in the range 60–120 km s⁻¹ (FWHM). The linear structure shows narrower line profiles in the range 48–87 km s⁻¹. Appendix A also shows that, although the majority of the gas in the ring has much higher radial velocities than the core structure, one small segment of the ring (at the point nearest the core) has a velocity similar to the core. This is supported by the CARMA spectrum, Figure 8(d), which shows a blueward component associated with the average spectrum of the ring. This strengthens the idea that the two CO structures are somehow related despite their discrepant radial velocities. We will discuss in the next section the possibility that the warm bridge connecting the CO ring and core may be splashed warm molecular gas caused by the collision of a dense molecular cloud moving at high (blueward) velocity with respect to more diffuse gas associated with material in the group.

Some of the CO emission in Field 5 may be associated with a larger filament and other CO clumps, as seen in a larger-scale (10''–15'' = 4.5–7 kpc) image, Figure 9(a), where we present the ALMA CO emission map smoothed (to 1 × 1 arcsec²) to bring out fainter features. This reveals a very faint filament of CO gas and several fainter southern CO concentrations. Both the filament and the southern CO structures have faint warm H₂ counterparts. Filaments like this one, in this case extending over 6''–10'' in scale (3–5 kpc), seem to be common in the MIRI maps of SQ (e.g., Figure 2). Figure 9(b) shows again the velocity field of the ring, compact core, and linear structure, this time in relation to the larger-scale structure in the region. This further emphasizes how the compact core has such a different velocity from all the other structures in the region. The velocity of the filament is poorly determined in CO because it is so faint, but seems to have a velocity similar to that of the average velocity of the group. Mid-IR spectroscopy would be needed to provide more information about the filament since it is well detected in the MIRI 10 μm image (presumed H₂ emission).

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4.1.3. Field 4: In the Northern Star-forming Region, SQ-A

Finally, we present the distribution of gas phases and star clusters in the Field 4 SQ-A region. The region has been discussed previously in terms of a possible tidal dwarf galaxy (Moles et al. 1997; Mendes de Oliveira et al. 2001), although this is disputed by Xu et al. (1999), who point out that the kinematics do not show simple rotation. Figure 10 shows a closer correspondence between the various gas phases, compared with Fields 5 and 6. Panels (a) and (b) of Figure 10 show a close correspondence between the CO distribution, the 10 μm H₂ image, and the 7.7 μm PAH-dominated image. In some of the panels, we have smoothed the CO distribution to an effective resolution of 1'' (458 pc). The close correspondence between these three interstellar medium (ISM) indicators is not unexpected since previous IRS spectroscopy of the region showed that the total H₂-to-PAH ratios in that region are more typical of normal photo dissociation region (PDR) regions associated with H ii regions. Also, ground-based spectroscopy shows that the optical emission-line ratios from the ionized gas are consistent with H ii-region-like spectra (Konstantopoulos et al. 2014). Extended patchy ionized gas is seen in panels (c) and (d) of Figure 10. A sharp boundary in the Hα distribution near the point where the CO emission starts to rise on the eastern edge, suggesting significant extinction of Hα light due to a dust lane associated with the denser molecular gas. This is supported by the fact that the same "edge" is not seen in the F200W NIRCam filter (Figure 10(d)), where the Paα is likely to be much less extinguished that the Hα. Konstantopoulos et al. (2014) measured high optical extinction in this region based on ground-based optical spectra. The emission seen in the F200W filter, Figure 10(d), shows many slightly extended clump sources scattered across an elongated "disk." The overall scale of this possible starburst dwarf system is 5 × 1.5 arcsec² (2.3 kpc × 700 pc). The U-band image supports the idea of strong extinction in the region of the peak CO emission. In Table B1, we also present UBV_mVI photometry of some of the brighter well-defined sources (Fedotov 2014). This study shows that the brighter clumps have ages based on the optical photometry ranging from 6 to 36 × 10⁷ yr, much longer than that of the blue clumps we investigated in Field 6. This may suggest a more mature stellar population, although because of higher extinction here NIRCam photometry would provide a more definitive result.

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The structure of the CO is that of two interleaved half-shells (very well seen in the smoothed version of the CO contours, which are shown in all the panels of Figure 10 except (b) and (d)). As with the previous descriptions, Figure 10(f) is an RGB composite of the F1000W (red, mainly H₂ dominated), the smoothed ALMA CO(2–1) (green, cool H₂), and the F200W (red, Paα plus starlight) images to show the spatial relationship between the different gas phases and star clusters.

The kinematics of the system led Mendes de Oliveira et al. (2001) to identify the region as a potential tidal dwarf because of its clear rotation using Fabry–Pérot observations with the Canada–France–Hawaii Telescope. They derived a systemic velocity of 6680 km s⁻¹, an asymmetric rotation curve, and a high-velocity dispersion of 115 km s⁻¹, which may have been the result of a blend of two different optical emission-line components.

Our CO spectra do not show obvious rotation. We show the integrated CO(2–1) emission map (moment 0) and velocity field map (moment 1) for Field 4 in Figure 11. Although the overall kinematics of the CO gas in Field 4 shows lower systemic velocity on the eastern side (6700 km s⁻¹) and a higher value to the west (6740 km s⁻¹), the velocity range is narrow (FWHM = 47 km s⁻¹), centered on the barycentric velocity of the main SQ group. The overall systemic velocity is in reasonable agreement with the optical emission-line velocities, although there is little real hint of rotation. This may argue against the formation of the object within a tidal stream, where the rotation is expected as a remnant of tidal shear. As we show in Table A3 and Figure A3, the velocity dispersion of the clumps at the highest resolved scale is narrow, which is very different from the more turbulent emission in Fields 5 and 6.

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Although not shown in Figure 10, there is a second very faint ALMA CO velocity component (6900 km s⁻¹), which is seen at a higher amplitude in CO(1–0) observed with the larger beam of CARMA data. It is very clear as a double-horned profile in the CARMA spectrum of Figure 11(c). This emission must be quite diffuse and is only just detected with ALMA. A similar high-velocity CO component was seen associated with SQ-A by Guillard et al. (2009) with the 30 m IRAM telescope. There is a hint of this higher-velocity component in the kinematic maps obtained in Hα by Duarte Puertas et al. (2019), where the ionized emission is quite localized to SQ-A. So, like Field 5, we find two CO components associated with a major CO concentration in SQ with quite discrepant velocities. In this case, the two components are roughly coincident and so its difficult to know whether the warm H₂ gas, which dominates the 10 μm JWST image, originates in two clouds in physical contact or are simply chance associations. Only mid-IR spectroscopy will allow us to fully explore this possibility.

In Table 1, we present measured and derived properties of the warm H₂ emission derived from both the Spitzer IRS data of Appleton et al. (2017) and from measured fluxes from the F1000W image for Fields 5 and 6, where we are confident that the emission is primarily from the 0–0 S(3) H₂ line. Unlike the tables in the Appendix A, which are concerned with illustrating the detailed kinematics and cold gas properties at the scale of the CO clumps, here we are concerned with the masses of the warm H₂ and the cold molecular mass integrated over the same larger areas, presented in Table 2. For Field 4, where we were unable to properly correct for underlying dust emission across the map, we only quote the properties of the H₂ from the IRS data on a large scale for comparison with Fields 5 and 6.

Table 2. Integrated Properties of CO Emission in Fields 5 and 6 over the Same Areas as Table 1

Region	Vhelio	FWHM	Σ(Sv ΔV)	Σ(Sv ΔV) a	M(H2) a	Mgas b
			CO(2–1)	CO(1–0)	×106	×106
	(km s−1)	(km s−1)	(Jy km s−1)	(Jy km s−1)	(M⊙)	(M⊙)
[1]	[2]	[3]	[4]	[5]	[6]	[7]
F5R1A	6401 (±6)	101 (±15)	0.9 (±0.1)	0.28 (±0.05)	18.8 (±3.0)	25.6 (±4.1)
F5R1B	6646 (±4)	73 (±10)	0.6 (±0.1)	0.19 (±0.03)	13.0 (±2.0)	17.7 (±2.7)
F5R2	6388 (±4)	119 (±10)	0.8 (±0.1)	0.25 (±0.02)	16.7 (±1.5)	22.7 (±2.0)
F5R3	6656 (±2)	105 (±5)	1.5 (±0.1)	0.47 (±0.02)	31.5 (±1.5)	42.8 (±2.0)
F5R4	6618 (±4)	84 (±8)	1.1 (±0.1)	0.33 (±0.04)	22.4 (±2.4)	30.4 (±3.3)
F5R5	6537 (±5)	67 (±12)	0.6 (±0.1)	0.20 (±0.04)	13.5 (±2.7)	18.3 (±3.6)
F6-ALL	6054 (±2)	91 (±4)	4.7 (±0.2)	1.5 (±0.1)	99.1 (±4.4)	134.8 (±6.0)
F6R1	6073 (±1)	47 (±2)	1.1 (±0.1)	0.3 (±0.0)	23.0 (±1.1)	31.3 (±1.4)
F6R2	6040 (±2)	93 (±4)	1.2 (±0.1)	0.4 (±0.0)	26.1 (±1.3)	35.4 (±1.7)
F6R3	6051 (±2)	41 (±5)	0.2 (±0.0)	0.1 (±0.0)	4.3 (±0.6)	5.9 (±0.8)
F6R4	6009 (±2)	37 (±4)	0.3 (±0.0)	0.1 (±0.0)	5.9 (±0.7)	8.1 (±1.0)
F6R5	6103 (±4)	48 (±10)	0.2 (±0.1)	0.1 (±0.0)	4.9 (±1.1)	6.7 (±1.5)

Notes.

^a We assume a conversion between a line flux at CO (1–0) to that of CO (2–1) of S_CO(1−0)/S_CO(2−1) = ${({\nu }_{1-0}/{\nu }_{2-1})}^{2}{({r}_{21})}^{-1}$, where ν₁₋₀ and ν₂₋₁ are the rest frequencies of the transitions, and r₂₁ is assumed to be 0.8, similar to that measured in the Taffy Galaxies bridge (Zhu et al. 2007), which shares many similarities with the gas in Stephan's Quintet (see Appleton et al. 2022). The H₂ masses presented here are for an X_CO value = X_CO,20 = N(H₂)/I_CO = 2 × 10²⁰ cm⁻² (K km s⁻¹)⁻¹, the standard value assumed for our Galaxy. In the text we discuss how this may not be applicable in such a turbulent region. M_H2 = 7.72 × 10³ D²Σ(S_v ΔV)(1+z)⁻¹, where D = 94 Mpc and Σ(S_v ΔV) is the CO₍₁₋₀₎ line integral with S_v in janskys and the velocity V of the gas in kilometers per second. We assume z = 0.02. ^b Total gas mass M_gas = M_H2 × 1.36, includes a 36% correction for helium (Bolatto et al. 2013).

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In Table 1, Column 1 gives the name of the regions defined in Figure 12 for Fields 5 and 6. The regions were chosen to cover the main warm H₂ features. Column 2 gives the MIRI filter used or the Spitzer module used (see later). For most of the regions, we extract the equivalent 0–0 S(3) flux from the JWST F1000W image. The pure rotational molecular hydrogen transition is listed in Column 3. Column 4 gives the measured H₂ flux density measured from the JWST F1000W over an area in arcsecs given in Column 5. We convert these flux densities to equivalent line fluxes, Column 7, using the filter bandwidth listed in Column 6. The line luminosity of the 0–0 S(3) line is given in Column 8, assuming a distance to SQ of 94 Mpc. We also present the properties of warm H₂ taken from Spitzer for the regions, and the appropriate IRS modules and H₂ rotational transition is listed. These are used below in the calculation of H₂ masses. Uncertainties are provided in parentheses for each measured quantity.

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To calculate the total warm H₂ mass in the regions, we make the following assumptions. First, since in most cases we only have the line flux of the 0–0 S(3) line derived from the MIRI image, we must assume something about the temperature distribution of the gas in order to estimate the total warm H₂ mass. Since the regions we are interested in were observed with Spitzer over a larger area, but over many of the rotational H₂ lines, we can assume, as an approximation, that the temperature distribution measured by the Spitzer IRS (Appleton et al. 2017) can be applied to all the smaller regions resolved by MIRI imaging. In that paper, we showed that the excitation diagrams for many regions extracted over SQ were almost always consistent with a single power-law relationship between the warm H₂ column density and the H₂ temperature of the form δN(H₂) ∝T⁻ⁿ δ T, following the work of Togi & Smith (2016). For Field 6, the Spitzer data provided a power-law index n = 4.5, which is quite shallow compared with normal H ii regions, and consistent with warm gas that is shock heated (Appleton et al. 2017).

Because Field 4 (the region in SQ-A) was ambiguous regarding isolating the H₂ line in the F1000W filter because of contamination from neighboring broad wings of PAH features (Figure 3), we did not attempt to measure the masses of smaller regions in this field. Rather we provide, for comparison with Fields 5 and 6, the IRS-measured values. This analysis yielded a much steeper value of n = 5. This is consistent with our observation that Field 4 is dominated by star formation.

Field 5 is a little more complicated. Unlike the other cases, this region lies outside the full spectral mapping region performed by Spitzer, and only the IRS Long-Low module was used to cover this area. This provided measurements of the 0–0 S(0) 28 μm and 0–0 S(1) 17 μm transitions, but not the shorter-wavelength transitions. However, the JWST observations from MIRI can provide the 0–0 S(3) flux, if we extract over the same extraction region (referred to as F5-IRSarea in Table 1) observed by Spitzer. With three points on the excitation diagram given in the table, we were able to measure a power-law index of 4.2. The flat value for n is comparable with some of the warmest regions of Stephan's Quintet (Appleton et al. 2017).

The above measurement assumes that the MIRI and Spitzer data can be compared reasonably accurately given that one is obtained from an image and the other from a spectrometer. To check the validity of this comparison, we compared the measured 0–0 S(3) line flux from IRS with the flux we derived from the extraction of the same area for Field 6. As Table 1 shows, we measured a line flux of 4.6 (±0.7) × 10⁻¹⁸ W m⁻² for the 0–0 S(3) line from the JWST image. With the IRS, we measured (region 142 of Appleton et al. 2017) a line flux of 4.7 (±0.3) × 10⁻¹⁸ W m⁻². The agreement provides some confidence that we are measuring the fluxes correctly using the MIRI F1000W image.

Having established an average power-law index, we now apply that same value of n appropriate to the smaller extracted regions for Fields 5 and 6 (Column 9 of the table). Without spectroscopic observations on the subarcsec scale, this is obviously an approximation. Using the 0–0 S(3) flux measured by JWST, we estimate the total H₂ mass in these regions integrated over the power law down to a temperature of 100 K. These masses (and uncertainties) are given in Column 10 of Table 1. We caution that these uncertainties do not include the unknown effect of possible variations, from region to region, in the actual power-law index, which we have assumed is equal to the average value derived over a much larger area.

From this analysis, we see that the warm H₂ mass of the entire arrow-head like structure in Field 5 (F5R1) is 5.6 (±1.0) × 10⁶ M_⊙ and its core (F6R2) contains almost half of the mass, 2.1 (±0.3) × 10⁶ M_⊙. The partial ring structure (F5R3) has a much smaller warm H₂ mass of 0.9 (±0.2) ×10⁶ M_⊙. The two more distant regions of warm H₂ gas at the end of the filament in Field 5 have larger warm gas masses of (F5R4) 2.5 (±0.4) and (F5R5) 1.8 (±0.3) × 10⁶ M_⊙, respectively.

In the case of the head-tail structure of Field 6, the triangular hot spot (F6R1) has a warm H₂ mass of 4 (± 0.7) × 10⁶ M_⊙ and most of the clumps in the tail have masses of 1–2 × 10⁶ M_⊙. The total H₂ mass integrated over the whole structure is 14 (± 3) × 10⁶ M_⊙, suggesting that there is considerable extended emission in the tail.

Fields 4, 5, and 6 have similar total warm H₂ masses when integrated over the larger 5.5 × 5.5 arcsec² (2.5 × 2.5 kpc) area measured by Spitzer IRS.

How do these warm H₂ masses compare with the mass of molecular gas measured by ALMA? In Table 2, we present the measured and derived properties of the gas from the ALMA observations. We use the integrated CO(2–1) line profile for the same regions as those used in Table 1 for the warm H₂. Here, we are concerned with measuring the total CO line profile associated with each of the regions for which we measure warm H₂ emission. Column 1 gives the region description. Columns 2 and 3 provide the mean heliocentric velocity and FWHM of the emission after fitting each region with a Gaussian line profile. We note that for F5R1 there we two detected CO emission features, referred to as F5R1A and B. This is because R1 contains emission from both the compact core feature and the linear filament. The inclusion of two distinctly different structures (with different kinematics) within the one region is clearly seen in Figure 12(b). These structures will be discussed in more detail in Appendix A. Column 4 provides the line integral of the CO(2–1) emission in jansky-kilometer per second, and Column 5 provides an estimate of the equivalent CO(1–0) emission, based on an assumed ratio of the integrated intensity of the main-beam temperature in the two lines (r₂₁ = I[12CO(2–1)]/I[12CO(1–0)]). Values of r₂₁ range from 0.76 to 0.9 in normal galaxies at low and intermediate redshift, with a slightly larger range for perturbed galaxies (Braine & Combes 1992; Daddi et al. 2015). We assumed a value of 0.8 to be similar to that measured in the highly turbulent Taffy Bridge (Zhu et al. 2007). Column 6 provides an estimate of the H₂ mass for an assumed value of the X-factor, X_CO. For the purposes of this discussion, we present the H₂ masses in Column 6 and the total gas mass in Column 7, assuming X_CO is the average Galactic value. This value is likely to be significantly overestimated in regions of strong shocks and turbulence. The footnotes to the table provide more details about the definitions and assumptions used in the calculations.

Bearing in mind the uncertainties in both the warm and cold molecular masses, we can estimate the ratio of M(H₂)_warm/ M(H₂)_cold by comparing the warm gas masses estimated by JWST with those from ALMA. For the F5R1 regions (the arrow-head structure in Field 5) the ratio of the warm-to-cold H₂ masses is 5.6/(25.6 + 17.7) = 0.13. This assumes that both the CO linear filament and the CO core both contribute to the total mass. It is not clear that the linear structure is actually detected in warm CO, so the ratio could be much higher if that object was excluded (0.22). However, in reality, since X_CO is likely to be much smaller than Galactic (say one-fifth X_CO,20), then the fraction of warm gas with T > 100 K is likely to be very large (almost equal to the cold H₂). These values are significantly greater than those seen in normal galaxies. F5R3, which includes the partial ring of CO emission, has a warm-to-cold H₂ ratio of 0.02, which is much more like a normal galaxy. Regions F5R4 and R5 show warm-to-cold ratios somewhat in between the two extremes (0.08 and 0.1). These regions lie at the opposite end of a warm filament (only weakly detected in CO) that connects them to the other Field 5 structures. The warm triangular hot spot in Field 6 shows an average ratio of warm-to-cold H₂ = 0.12. Again, if we assume X_CO is one-fifth of the Galactic value, this would imply a significant warm component. F6R2 has a much lower ratio of warm-to-cold gas (0.03), a number which is closer to that of normal star-forming galaxies. It is interesting that this is the region which emits strongly in the F770W filter, which is consistent with this containing a PDR. The other clumps in the tail have unusually large values of warm-to-cold CO (in the range 0.18–0.22), reflecting the relative faintness of the CO in the clumps in the tail of Field 6. The overall conclusion is that, except for the regions which show obvious star formation (PAH emission and compact star clusters embedded in the emission), the warm-to-cold H₂ ratios are unusually large for both Field 5 and Field 6. We will discuss this further in the discussion in the context of shocks and turbulence.

We note that our ALMA radio continuum observations at 230 GHz might, in principle, provide useful information about the mass of cool dust associated with the target regions. Inspection of the continuum maps did not detect dust emission from any of the regions of interest down to an rms upper limit of 0.03 mJy beam⁻¹. We assume a dust mass M_dust = S_ν D²/(κ_ν B(ν,T_dust)), where S_ν is the flux density at frequency ν, D is the distance to the galaxy, and B(ν,T_dust) is the blackbody function at frequency ν and temperature T_dust. Assuming a reasonable average value for the dust absorption coefficient κ_{1.3 mm} of 0.03 m² kg⁻¹ (extrapolating from 500 μm from Clark et al. 2019), T_dust = 20 K, we estimate a 3σ upper limit to the dust mass of M_dust < 5.6 × 10⁵ M_⊙ beam⁻¹. A dust temperature of 20 K is consistent with previous Herschel observations of dust emission on much larger scales (Appleton et al. 2013). If the typical gas-to-dust ratio for galaxies (Sandstrom et al. 2013) is 80–100 (this number may be higher if dust is destroyed in the shock waves), the corresponding upper limit to the gas mass is at least a factor of 10 larger than the CO-derived gas masses within the same beam. This suggests that the continuum observations are not deep enough to provide useful physical constraints on the cold dust properties of the observed regions. Observations at higher frequencies would be desirable to potentially detect cold dust from these regions.

In order to explain the formation of large masses of molecular gas and high luminosity in the pure rotational ground-state transitions of H₂, Guillard et al. (2009) presented a model in which a large-scale driving shock wave (traveling at 600–800 km s⁻¹) from the intruder passes into a multiphase medium of gas in a filament of tidal debris caused by previous tidal interactions between group members (see Figure 13 for a schematic view). We summarize here that scenario, which was based on observations with limited spatial coverage and resolution with Spitzer, and then discuss what we are now learning from the new JWST plus ALMA data presented here about the link between the large-scale dynamics in the group and the formation and excitation of the molecular gas in the IGM.

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The large-scale shock wave driven by the intruder, NGC 7318b, passes over the tidal, multiphase filament, heating low-density gas (n_H < 2 × 10⁻² cm⁻³) to X-ray-emitting temperatures (see O'Sullivan et al. 2009), and destroying any dust grains present there (illustrated in Figure 13 as the light orange medium). Regions in the pre-shock region with overdensities greater than 10 (typically H i gas with n_H > 0.1 cm⁻³) compared to the volume-filling gas are heated to lower (< 10⁶ K) temperatures, and can cool and form molecules on dust grains as they survive the slower shocks (sketched as the dark blue clouds in Figure 13). Dust grains are expected to survive in shocks with velocities lower than 200 km s⁻¹ (Jones & Nuth 2011). Guillard et al. (2010) have indeed shown that dust survives in the denser regions of the IGM of the galaxy collision, and that the PAH and continuum IR emission is consistent with the dust associated with the warm (>150 K) H₂ gas, heated by a UV radiation field of intensity comparable to that of the solar neighborhood, G₀ = 1.4. Natale et al. (2010) also showed that a diffuse IR dust emission component is associated with the collisional cooling of the hot halo gas, provided that dust is injected into the hot plasma by tidal stripping or stellar feedback. In the next section, we discuss how the dynamical interaction between the cold and hot gas can lead to the same physical conditions downstream from the shock, as thermally unstable gas is produced within the turbulent mixing layers (Gronke & Oh 2022), leading to a cycle of molecular clouds breaking up and reforming in the flow. This is sketched as the turbulent tail in the zoom-in panel on the right-hand side of Figure 13.

An efficient transfer of the bulk kinetic energy of the galaxy collision to turbulent energy within the molecular gas, associated with a continuous cold gas destruction–reformation cycle, is required to make H₂ a dominant coolant of the post-shock gas (Guillard et al. 2009). In the next section, we discuss how the spatial decoupling between cold and warm gas at GMC scales, revealed by the comparison between ALMA and JWST images, is shedding light on the physical processes that drive this energy transfer mediated by thermal instability and a dynamical interaction between gas phases.

The ALMA and JWST data presented in this paper reveal complex and turbulent structures, with large velocity dispersions (40–100 km s⁻¹) at GMC scales, and a spatial decoupling of cold and warm molecular gas. Although a quantitative interpretation of these data is difficult and beyond the scope of this paper, we qualitatively discuss a scenario that relate the cold and warm H₂ distributions to theoretical work of cold clouds shattering in their dynamical interaction with the hot volume-filling hot plasma and subsequent gas cooling (Gronke & Oh 2020).

When the shock wave hits the multiphase IGM filament, the clouds are compressed on a crushing timescale, t_crush = L_c /V_c , where V_c is the shock velocity in the cloud and L_c the size of the pre-shock cloud. The shocked clouds experience dynamical instabilities and mixing at the interface with the hot gas flow, as well as thermal instabilities induced by gas cooling, which provide an additional fragmentation mechanism occurring over the gas cooling timescale, t_cool, as the shock moves into the cloud. The gas stripped from the cloud may cool and condense as a result of thermal instability. This occurs if the gas cooling time t_cool is smaller than t_crush. There is a wide range of parameter space (typically n_H ≈ 0.1–100 cm⁻³, L_c = 1−100 pc) where t_crush > t_cool, i.e., molecular gas can form before the shock has crossed the cloud and thus before dynamical instabilities fully develop (see Figure 4 in Guillard et al. 2009). The cooling time in the warm molecular cloudlets is mostly driven by H₂ cooling, therefore the dust-to-gas mass ratio could, in principle, indirectly affect the H₂ reformation in the post-shock gas, and therefore the cooling time. However, modeling of the dust-to-gas ratio, and the the PAHs to very small grains abundance ratio, based on Spitzer and Herschel spectroscopy and photometry, show that these are comparable to the Galactic value in the molecular gas (Guillard et al. 2010). This again suggests that dust survives the slower shocks driven by the dynamical instabilities and mixing into the molecular cloudlets.

Because of thermal instability, the post-shock clouds shatter in many clumplets, which then are seeds for further cooling-induced condensation, but also mixing (Farber & Gronke 2022a). In this parameter space where the cooling time is shorter than the cloud destruction timescale, this shattering process causes the gas to cycle through the ISM phases with the reformation of cold clumpy gas. Cloud growth depends on density as well as on local heating processes (Jennings et al. 2023). Given the large masses of diffuse molecular gas observed in the IGM where atomic H i gas is undetected, qualitatively, the post-shock physical conditions must be in the density and irradiation regime where the cloud evolution is dominated by shattering and perhaps cooling-driven coagulation.

Many of the structures detected with JWST along the main north–south molecular filament are shaped like cometary tails, with Field 6 being a prime example (see Figure 12). They could be signatures of a misty H₂ outflow driven by the dynamical interaction between cold and hot ISM phases, e.g., ram pressure stripping of CO clouds as the hot X-ray-emitting gas flows past them. This is illustrated in the zoom-in panel of Figure 13. The high abundance of warm H₂ relative to the CO-derived cold gas mass discussed in Section 4.1 for both Field 6 and Field 5 lends support to this idea, despite their different spatial positions in the SQ group. We will return to these differences later. The pressure-driven flow velocity is larger than the turbulent gas velocity because the droplets are tiny. At the scale of an H₂ droplet, the difference between the turbulent velocity dispersion and the flow velocity is even larger. Farber & Gronke (2022a) suggest that in typical circumgalactic medium or galaxy cluster environments, cloud shattering during the cooling process could be responsible for in situ formation of a fog of warm (1000 K) gas embedded in hot (10^7–8 K) plasma, with molecular shattering length scales of the order of 0.1–100 pc. This would imply that fragments formed in such manner are very small, with a low volume-filling fraction. A detailed investigation of the physical conditions, which could be achieved with JWST H₂ spectroscopy, would allow us to determine whether the cloudlets can coagulate or whether the molecular streams we observe remain misty.

The very complex, multiscale kinematical and morphological structures observed at high resolution with JWST and ALMA points toward a highly clumpy (misty) structure of the molecular gas in the IGM of SQ. This could explain why the Lyα line emission is so strong and the line profile so broad (Guillard et al. 2022). Indeed, the cloudlets must have a small volume-filling fraction but their surface filling factor must be close to unity so Lyα photons can escape through multiple scatterings off the clump surfaces without being absorbed in the inter-clump, dust-free plasma.

The ALMA and JWST observations also shed light on a puzzle raised by the H₂ observations by Spitzer (Appleton et al. 2006), as well as single-dish CO observations, which showed extreme velocity dispersion (≈1000 km s⁻¹) of the molecular gas on kiloparsec scales. These large-scale velocity gradients, also observed in other gas phases, could partly originate from coherent bulk flows (Guillard et al. 2022). The H₂ molecules can only survive at shock velocities ≲20 km s⁻¹, which is much smaller than those velocities at kiloparsec scales, but also smaller than the line widths of the CO clouds. The ALMA observations presented in this paper indeed show that the CO clumps have velocity dispersions in the range of 40–100 km s⁻¹. Those velocities are also higher than the shock velocities needed to account for the excitation of the pure rotational lines of H₂ of the order of 5–20 km s⁻¹ (Guillard et al. 2009; Lesaffre et al. 2013; Appleton et al. 2017).

The tiny H₂ droplets are unlikely to survive over the length of the tails, which supports the idea that they are continuously destroyed and reformed out of the more diffuse gas. This could explain why some of the CO clumps (in the tail of Field 6, for instance) have high-velocity dispersions (see Figure A1), as they are reformed out of pressure-driven outflow from the parent cloud. This will need to be further investigated with future H₂ spectroscopy with JWST to study the velocity gradients of the warm H₂ gas at the scales of the ALMA observations. Spectroscopy would also allow us to see whether this fog of small warm clouds is being accelerated in the bulk flow of the hot X-ray-emitting gas, and would likely differ from the motion of the massive CO clouds. The temperature of the warm droplets may also increase along the tail as the foggy clouds progressively mix and become heated in the faster flow. This might be evident in the excitation diagram of H₂ when a full complement of H₂ rotational lines can be measured. Those observations will motivate a direct comparison with future numerical simulations.

The comparison between the warm and cold molecular gas spatial distributions in the IGM of SQ shows that the powerful mid-IR emission originally discovered with Spitzer spectroscopy does not primarily have a one-to-one correlation with the cold CO(1–0)-emitting gas. This supports the view of the model put forward in Guillard et al. (2009), which involves a turbulent energy cascade and dissipation of mechanical energy within the warm H₂ phase, both driven by the dynamical interaction between the hot and cold gas phases. The complex morphology and kinematics of the warm and cold H₂ structures shows that this dynamical interaction occurs over a wide range of scales. This is expected from an energy cascade over many orders of magnitude in spatial scales, from tens of kiloparsecs at the injection scale of mechanical energy from the galaxy collision to subparsec dissipation scale through line emission (Guillard et al. 2009).

As we have stressed earlier, the high luminosity in the warm molecular hydrogen lines over a large part of the Stephan's Quintet system requires a common mechanism to convert a fraction of the kinetic energy resulting from the large-scale shock wave into rotational molecular hydrogen via low-velocity shocks (Guillard et al. 2009; Appleton et al. 2017). We have seen that, of the three regions we studied in detail here, the one that most closely fits the proposed cloudy outflow model is the head-tail structure in Field 6, where we seem to be witnessing a possible breakup of a large cold molecular cloud into warmer H₂ in both the head and tail.

Although modeling of the actual structure of a cold molecular cloud in a hot wind will require further simulation, it is likely that cometary structures should be common along the main shock wave, and the diversity of structures should be similar to those seen in the cloud-survival sequences of models by Farber & Gronke (2022b). Those model sequences, where clouds are entrained in a fast wind, show that in some cases the cooler head and tail structures can survive for long periods, whereas in other cases, depending on the properties of the clouds, they can dissolve fairly quickly. Indeed, one of the models has a morphology remarkably similar to the shell-like CO and warm H₂ structure in the head of Field 6, including the spike feature at the head of the clump. Many of the shapes of the bright H₂-emitting clumps seen along the main shock front in Stephan's Quintet (Figure 2) show suggestions of head/tail structures, and have filamentary structure very reminiscent of the modeled shapes seen when the head/tail structures survive. Others have shapes similar to the models where the clouds have lost their heads and are in the process of dissolving. Therefore, it is plausible that our picture of the foggy outflowing cloud picture is applicable to a large number of cloud structures seen in the main shock.

In addition, Figure 2 shows head-tail structures far to the south and west of the main SQ-A region. One of those tails extends over 30''–40'' (10–20 kpc!). These structures may represent the interaction of dense molecular clumps with the outer hot halo of the intruder far downstream of the main shock. More observations of the cold and warm molecular gas and their kinematics will be needed to test this idea.

In the case of our study of the Field 4 region in SQ-A, it is less clear that the cloud outflow model is applicable, although we note that a faint irregular tail is seen extending to the east of the main concentration of gas and stars. Xu et al. (1999) have previously suggested that SQ-A is too young to have formed as a tidal dwarf galaxy, but is more likely a starburst region created as part of the main shocked filament. The H₂ gas temperature measured by Spitzer in SQ-A is much lower than other regions in the main filament (Appleton et al. 2017), and the gas may have cooled to the point that stars can form in large quantities there. Although understanding how those stars form is beyond the scope of this paper, we have shown in Guillard et al. (2009) that the large (>50 pc), dense (n_H > 100 cm⁻³) pre-shock clouds are expected to be gravitationally unstable in the hot gas at this pressure (see their Figure 4). Therefore, their evolution will be different than the more diffuse gas, and should trigger star formation. This may explain star formation in SQ-A and other star-forming regions scattered in the IGM. These large pre-existing clouds have crushing times longer than the SQ collision age (believed to be a few tens of millions of years). However, such clouds must be rare, otherwise we would see major bursts of star formation all over the main filament, which is not observed.

The case of Field 5 (in the bridge between NGC 7319 and the main filament) is interesting. As we have shown in Section 4, the morphology of the warm and cold gas is not so obviously a head-tail structure, although there is a warm gas filament running through the region. This region may owe its overall morphology and kinematics to another process—for example, the collision of a dense molecular concentration (the compact core) with a diffuse region near NGC 7319, driving a ring into the target material. The elongation of the compact CO-emitting core along the direction which points toward the center of the CO-emitting ring, and the existence of a clump of CO in the ring at the same velocity as the core, may support a dynamical connection. Future warm H₂ spectroscopy will determine whether the warm gas connecting the two CO structures with disparate radial velocities is related, or merely chance alignments of two different cloud systems along the line of sight. Such events, if they occur, are likely to be relatively rare, and are unlikely to provide a generic mechanism for the dissipation of collisional kinetic energy into widespread warm H₂ emission.

Finally, we come to the complex network of filaments around NGC 7319, which, for the most part, appear pink in the false-color representation of Figure 2. This implies that the 7.7 μm and 15 μm emission is much stronger here than in the regions which are warm H₂ dominated.

The irregular radial structure of the filaments might suggest a nuclear origin connected with the active galactic nucleus (AGN) in NGC 7319. The JWST MRS data in the central region of the AGN indeed reveal a spatially resolved ionized and molecular outflow, with kinematical evidence of an interaction between the AGN jet and the host ISM (Pereira-Santaella et al. 2022). Perhaps previous jet activity from the AGN has somehow formed these strange filaments.